The Advanced CCD Imaging Spectrometer (ACIS) offers the capability
to simultaneously acquire high-resolution images and moderate
resolution spectra. The instrument can also be used in conjunction
with the High Energy Transmission Grating (HETG) or Low Energy
Transmission Grating (LETG) to obtain higher resolution spectra (see
Chapters 8 and 9). ACIS contains 10
planar, 1024×1024 pixel CCDs (Figure 6.1); four
arranged in a 2×2 array (ACIS-I) used for imaging, and six arranged in
a 1×6 array (ACIS-S) used either for imaging or for a grating spectrum
read-out.
Two CCDs are back-illuminated (BI) and eight are front-illuminated
(FI). The response of the BI devices extends to energies below
that accessible to the FI chips. The chip-average energy resolution
of the BI devices is better than that of the FI devices.

In principle, any combination of up to 6 CCDs can be operated
simultaneously. However, because of changes in the thermal
environment, the CXC now recommends that fewer CCDs be selected
if the science needs can be met with fewer CCDs. Some
CCDs can be designated as optional, which means they may be
turned off depending on thermal conditions.
See the Section 6.21.1
for more details on CCD selection recommendations and policies.

Figure 6.1: A
schematic drawing of the ACIS focal plane; insight to the terminology
is given in the lower left. Note the aimpoints: on S3 (the
'+') and on I3 (the 'x').
Note the differences in the orientation of the I and S chips, important
when using subarrays (Section 6.12.1). Note also the (Y,
Z) coordinate system and the target offset convention (see
Chapter 3) as well as the SIM motion
(+/−Z).
This view is
along the optical axis, from the sky toward the detectors,
(-X). The numerous ways to refer to a particular CCD are indicated:
chip letter+number, chip serial number, and ACIS chip number (CCD_ID);
see also Table 6.1. As indicated, S3 and S1 are
back-illuminated (BI) CCDs, and the rest are front-illuminated (FI)
CCDs. The node numbering scheme and the row/column directions
are illustrated lower center.

CCD_ID

CCD Name

CCD Type*

Serial Number

0

I0, ACIS-I0

FI

w203c4r

1

I1, ACIS-I1

FI

w193c2

2

I2, ACIS-I2

FI

w158c4r

3f

I3, ACIS-I3

FI

w215c2r

4

S0, ACIS-S0

FI

w168c4r

5

S1, ACIS-S1

BI

w140c4r

6

S2, ACIS-S2

FI

w182c4r

7f

S3, ACIS-S3

BI

w134c4r

8

S4, ACIS-S4

FI

w457c4

9

S5, ACIS-S5

FI

w201c3r

fcontains I-array aimpoint; best imaging

fcontains S-array aimpoint; best imaging

*FI = front-illuminated, BI = back-illuminated

Table 6.1: Naming conventions for the ACIS CCDs.
See Figure 6.1 for the positions and orientations
in the flight focal plane.

Figure 6.2: A schematic drawing of the ACIS focal
plane, not to scale. The ACIS-I array consists of
chips I0-I3 (shaded gray in the upper figure).
The ACIS-S array consists of chips S0-S5 (shaded gray
in the lower figure).
See the discussion in Section 6.21.1
for information on selecting CCDs.

The original Instrument Principal Investigator for ACIS is Prof. Gordon Garmire
(Pennsylvania State University). ACIS was developed by a collaboration
between Penn State, the MIT Kavli Institute for Astrophysics and Space Research and the Jet
Propulsion Laboratory, and was built by Lockheed Martin and MIT. The
MIT effort was led by Dr. George Ricker. The CCDs were developed by
MIT 's Lincoln Laboratory.

ACIS is a complex instrument having many different
characteristics and operating modes. Radiation damage suffered by the
FI chips has had a negative impact on their energy resolution - the
BI devices were not affected - thus affecting the basic considerations
as to how to make best use of the instrument. We discuss the trade-offs in this
chapter. Software methods for improving the energy resolution of the
FI CCDs are discussed in Section 6.7.1. The low energy
response of ACIS has also been affected by the buildup of a
contaminant on the optical blocking filters and this is discussed in
Section 6.5.1.

Many of the characteristics of the ACIS instrument are summarized in
Table 6.2.

A CCD is a solid-state electronic device composed primarily of
silicon. A "gate" structure on one surface defines the pixel
boundaries by alternating voltages on three electrodes spanning a
pixel. The silicon in the active (depletion) region (the region below
the gates wherein most of the absorption takes place) has an applied
electric field so that charge moves quickly to the gate surface. The
gates allow confined charge to be passed down a "bucket brigade"
(the buried channel) of pixels in parallel to a serial read-out
at one edge by appropriately varying ("clocking") the voltages in
the gates.

The ACIS front-illuminated CCDs have the gate structures facing the
incident X-ray beam. Two of the chips on the ACIS-S array (S1 and S3)
have had the back sides of the chips treated, removing
insensitive, undepleted bulk silicon material and leaving the
photo-sensitive depletion region exposed. These are the BI chips and
are deployed with the back side facing the HRMA.

Photoelectric absorption of an
X-ray photon in silicon results in the liberation of a proportional number of
electrons: an average of one electron-hole pair for each 3.7 eV of energy absorbed. Immediately after the photoelectric
interaction, the charge is confined by electric fields to a small
volume near the interaction site. Charge in a FI device can also be
liberated below the depletion region, in an inactive substrate, from
where it diffuses into the depletion region. This charge may easily
appear in two or more pixels.

Good spectral resolution depends upon
an accurate determination of the total charge deposited by a single
photon. This in turn depends upon the fraction of charge collected,
the fraction of charge lost in transfer from pixel to pixel during
read-out, and the ability of the read-out amplifiers to measure the
charge. Spectral resolution also depends on read noise and the
off-chip analog processing electronics. The ACIS CCDs have read-out
noise less than 2 electrons RMS. Total system noise for the 40 ACIS signal
chains (4 nodes/CCD) ranges from 2 to 3 electrons (RMS) and is
dominated by the off-chip analog processing electronics.

The CCDs have an "active" or imaging section (see
Figure 6.1) which is exposed to the incident
radiation and a shielded "frame store" region. A typical mode of the
ACIS CCD operation is: (1) the active region is exposed for a fixed
amount of time (full frame ∼ 3.2 s); (2) at the end of the
exposure, the charge in the active region is quickly ( ∼ 41 ms)
transferred in parallel into the frame store; (3) the next exposure begins; (4)
simultaneously, the data in the frame store region is transferred serially to a
local processor which, after removing bias
(see Section 6.13), identifies the
position and amplitude of any "events" according to a number of
criteria depending on the precise operating mode. These criteria
always require a local maximum in the charge distribution above the
event threshold (see Table 6.2). The position and the
amount of charge collected, together with similar data for a limited
region containing and surrounding the pixel are classified
("graded") and then passed into the telemetry stream.

Since the CCDs are sensitive to optical as well as X-ray photons, optical
blocking filters (OBFs) are placed just over the CCDs between
the chips and the HRMA. The filters are composed of polyimide (a
polycarbonate plastic) sandwiched between two thin layers of
aluminum. The nominal thicknesses of these components for the two
arrays are given in Table 6.3. Details of the calibration
of these filters may be found in the ACIS calibration report at
http://www.astro.psu.edu/xray/docs/cal_report/node188.html.
These calibrations do not include the more recent effects of molecular
contamination. This is discussed in Section 6.5.1.

The brightness threshold for optical contamination (here defined as
a 1 ADU shift
per pixel in the signal) is based on on-orbit calibration observations
of two stars: Vega and Betelgeuse. The ratio of the signals leads us
to believe that there is a light leak in the I or J band in the near
infrared part of the spectrum. The back-illuminated (S3 and S1) chips
are much more sensitive, and we find threshold magnitudes of I ∼ 5.
The I-array is less sensitive by a factor of
about 50, for a limiting magnitude of I ∼ 0.75. Brighter
stars can be observed, but require special techniques in either data
analysis, observation planning, or both. Details can be found on this
web page:
http://cxc.harvard.edu/cal/Hrma/UvIrPSF.html. The CXC HelpDesk (http://cxc.harvard.edu/helpdesk/)
can help with detailed questions.

Measurements of ACIS include
laboratory calibrations, a system-level ground calibration of the
HRMA and ACIS at the X-Ray Calibration Facility (XRCF) at MSFC, and
on-orbit calibration using celestial and on-board radioactive X-ray
sources. The ACIS external calibration source (ECS) consists of
an 55Fe source
and a target made of aluminum and titanium. The source emits five strong
lines (Al Kα at 1.49 keV, Ti Kα and Kβ, at 4.51 and
4.93 keV, and Mn Kα and Kβ at 5.90 and 6.49 keV). A number
of weaker lines are also present.

The on-orbit calibration of ACIS is a continuing activity. All
calibration data are, or will be, described in detail, at
http://cxc.harvard.edu/cal/.
Consult the WWW site and its links for the
latest information.

The quantum efficiencies near the read-out for the ACIS CCDs for the
standard grade set, including optical blocking filters and
molecular contamination,
are shown in Figure 6.3. Note that the quantum
efficiency for the FI chips varies somewhat with row number (not
shown), and decreases by 5-15% farthest from the read-out at energies
above about 4 keV. This is due to the migration of good grades to bad
grades produced by charge transfer inefficiency (CTI), which varies with row
number. The quantum efficiency (QE) variation with position for the BI chips is much
smaller.

Cosmic rays tend to cause large blooms on FI chips, and much smaller ones
on BI chips. This results in a 2-4% decrement of the QE for FI chips
and ∼ 0.5% for BI chips.

The combined HRMA/ACIS on-axis effective areas are shown in
Figure 6.4 (log energy scale) and
6.5 (linear energy scale).
The effective areas are for an
on-axis point source and a 20′′-diameter extraction region. The ACIS effective areas include a correction that accounts for the
buildup of molecular contamination on the ACIS filters (see the discussion
in Section 6.5.1). Figures 6.4
and 6.5 show the predicted ACIS effective areas
for the middle of Cycle based on the current
time-dependent ACIS contamination models.

Figure 6.6 shows the vignetting
(defined as the ratio of off-axis to on-axis
effective area) as a function of energy at several off-axis angles. These
data are from a calibration observation of G21.5-0.9, a bright supernova
remnant/pulsar wind nebula.
The detector was appropriately offset for each off-axis angle so that the
data were obtained at the same focal position, minimizing the effects of any
spatially-dependent variations in the CCD response.

Figure 6.3: The quantum
efficiency (convolved with the transmission of the appropriate optical
blocking filter) of the FI CCDs (from a row nearest the read-out)
and the two BI CCDs as a function of energy. S3 is somewhat
thicker, hence more efficient, than S1.
These curves include the effects of molecular contamination, as
discussed in the text.

Figure 6.4: The
HRMA/ACIS predicted effective area versus the energy on a log
scale. The dashed line is for the FI CCD I3, and the solid line is
for the BI CCD S3.
These curves include the effects of molecular contamination, as discussed in the text.

Figure 6.5: The HRMA/ACIS predicted effective area versus the
energy on a linear scale. The dashed line is for the
FI CCD I3, and the solid line is for the BI CCD S3.
These curves include the effects of molecular contamination, as discussed in the text.

Figure 6.6: Vignetting (the ratio of off-axis to on-axis effective area) as
a function of energy for several off-axis angles in arcmin.

Astronomical observations and data acquired from the on-board ACIS external
calibration source (ECS) show that the ACIS effective area below 2 keV has
continuously declined since launch due to the build-up of out-gassed
material on the cold ACIS optical blocking filters. The HRC operates at a
warmer temperature and shows no sign of contamination build-up. The
build-up of contaminant on the ACIS filters has been monitored
with ECS observations, LETG/ACIS-S observations of the blazars
Mkn421 and PKS2155-304 and the isolated
neutron star RXJ1856-3754, and ACIS imaging observations of the rich
cluster of galaxies Abell 1795 and the oxygen-rich supernova
remnant E0102-72.3. The ECS is an 55Fe source with a
half-life of 2.7 years and is no longer useful for
monitoring the contamination on the ACIS filter.

Typically the calibration team releases updates to the contamination
model near the end of each calender year. The update accounts for any
changes in the characteristics of the contamination as measured
by observations taken over the same calender year. Most updates do
not affect the analysis of data taken prior to that calender year.
There are three components to the ACIS contamination model:
1) the build-up rate of the contaminant on the ACIS filters,
2) the spatial distribution of the contaminant on the ACIS filters, and
3) the chemical composition of the contaminant and hence its
energy-dependent effects. The energy-dependent effect is largest near
the Carbon edge at 0.277 keV and is negligible
beyond ∼ 1-2 keV.
Figures 6.7-6.9
show that all three of
these properties concerning
the contaminant have changed significantly during the Chandra mission.

Figure 6.7 shows that the optical
depth of the contaminant appeared to be
leveling off prior to about 2009, at which time the condensation rate of
the contaminant onto the ACIS filter began increasing. Also, before 2014,
the depth of the contaminant near the ACIS-S and ACIS-I aim-points was
approximately the same. The 2014 observations of A1795 show that the depth
of the contaminant is now greater on the ACIS-I filter compared to the
ACIS-S filter.

In addition, the contaminant has always been thicker near
the edges of the filters. Figure 6.8
shows that the difference in the
optical depth on the filter measured near the center of the chip
and near the edge of the chip appeared to be leveling
off prior to 2009, but has been increasing since this time.
The data for the edge of the chip are taken at CHIPY=64 (near the
readout) for ACIS-S (and equivalent values for ACIS-I).

LETG/ACIS-S observations show that the contaminant is composed of
carbon, oxygen and
fluorine. Figure 6.9 shows that the
oxygen-to-carbon ratio of the material
condensing onto the ACIS filters has varied significantly during the mission
indicating the presence of several different contaminants which have
dominated the condensation rate onto the ACIS filters at different times.

To assist with the proposal process, PIMMS tables are generated for
the ACIS effective area based on an extrapolation of the current
contamination model to the middle of the next Cycle, which is
an extrapolation of approximately 18 months. Based on recent changes
in the behavior of the out-gassed material condensing onto the ACIS filters,
proposers are advised to include an uncertainty of about 10% in
the predicted ACIS on-axis count rates below about 1 keV, and
about 20% near the chip edge.

Figure 6.8:
Difference in the optical depth between near the center of the chip
and the edge of the chip near the readout versus time.
ACIS-I3 data are shown as light gray (red) circles, and
ACIS-S3 data as dark gray (blue) circles.
The black circles are ECS data.
(See the online Proposers' Observatory Guide for a color version of this figure.)

Figure 6.9:
The ratio of the deposition rate in oxygen
atoms per year to that for carbon over the Chandra mission as measured
by LETG/ACIS-S observations of blazars.

The spatial resolution for on-axis imaging with HRMA/ACIS is limited by
the physical size of the CCD pixels (24.0 μm square ∼ 0.492
arcsec) and not the HRMA. This limitation applies regardless of
whether the aimpoint is selected to be the default aimpoint on
I3 or S3
(Figure 6.1). Approximately 90% of the encircled
energy lies within 4 pixels (2 arcsec) of the center pixel at 1.49 keV
and within 5 pixels (2.5 arcsec) at 6.4 keV. Figure 6.10
shows an in-flight calibration. There is no evidence for any
differences in data taken with either S3 or I3 at the nominal focus.
The ACIS encircled energy as a function of off-axis angle is
discussed in Chapter 4 (see Section 4.2.2
and Figure 4.13).

Off-axis, the departure of the CCD layout from the ideal focal
surface and the increase of the HRMA PSF with off-axis angle become
dominating factors. Since the ideal focal surface depends on energy,
observers for whom such considerations may be important are urged to
make use of the MARX simulator to study the impact on their
observation.

Figure 6.10:
The on-orbit encircled broad-band energy versus radius for an
ACIS observation of point source PG1634-706 (ObsID 1269).
The curve is normalized to unity at a radius of 5 arcsec. We estimate
that the statistical uncertainty is 3%, and the systematic uncertainty
due to power beyond 5′′ is 2%. The effective energy is 1 keV.

The Chandra PSF near the aimpoint displays an unexplained enhancement
in the profile ≈ 0.8" from the source centroid
(see Section 4.2.3). This anomaly is in
excess of that expected from ray trace simulations, and is preferentially
oriented towards the mirror spherical coordinate (MSC) angle
of φ = 285°
(see the CIAO caveats page
http://cxc.harvard.edu/ciao/caveats/psf_artifact.html).
This is approximately oriented towards the spacecraft +Z axis (see
Figure 1.2). The asymmetry is illustrated for
an ACIS observation of an LMXB in NGC 6397 in
Figure 4.17, and the magnitude of the asymmetry is
illustrated for a number of low-count-rate on-axis point sources in
Figure 4.18 (see
Section 4.2.3).
Figure 6.11 depicts the effect of the
anomaly on a number of point sources in NGC 6397.

Figure 6.11: The PSF anomaly illustrated with an ACIS-S observation of
NGC 6397 for a number of sources near the aimpoint.
The image has been derolled such that the spacecraft +Z is
pointed downwards (see solid vertical downward arrow of length 5").
The nominal direction of the anomaly is indicated by
the dashed arrow (also of length 5"). The data have been binned
to 1/4 ×1/4 ACIS sky pixels, and a logarithmic intensity stretch
has been applied.
Notice that all the sources in the field have discernible
asymmetries in the indicated direction.

The ACIS FI CCDs originally approached the theoretical limit for the
energy resolution at almost all energies, while the BI CCDs exhibited
poorer resolution. The pre-launch energy resolution as a function of
energy is shown in Figure 6.12. Subsequent to launch
and orbital activation, the energy resolution of the FI CCDs has
become a function of the row number, being near pre-launch values
close to the frame store region and substantially degraded in the
farthest row. An illustration of the dependence on row is shown in
Figure 6.13.

Figure 6.12: The
ACIS pre-launch energy resolution as a function of energy.

The loss of energy resolution is due to increased charge transfer
inefficiency (CTI) caused by low energy protons encountered during
radiation belt passages and Rutherford scattering through the
X-ray telescope onto the focal plane. Subsequent to the discovery
of the degradation, operational procedures were changed:
ACIS is moved to a sheltered position
during radiation belt passages. Since
this procedure was initiated, no further degradation in performance
has been encountered beyond that predicted from pre-launch
models. The BI CCDs were not impacted, which is consistent
with the proton-damage scenario - it is far more difficult for
low-energy protons from the direction of the HRMA to deposit their
energy in the buried channels of the BI devices, since the channels
are near the gates and the gates face in the direction opposite to the
HRMA. Thus the energy resolution for the two BI devices remains
nearly at pre-launch values.
The position-dependent energy resolution of the FI chips
depends significantly on the ACIS operating temperature. Since orbital
activation, the ACIS operating temperature has been lowered in steps
and is now set at the lowest temperature thought safely and
consistently achievable ( ∼ −120°C).

The ACIS instrument team has developed a correction algorithm that
recovers much of the lost energy resolution the FI CCDs.
The correction recovers a significant fraction of the CTI-induced
loss of spectral resolution of the FI CCDs at all energies.
The algorithm has been incorporated in the CIAO tool
acis_process_events since CIAO 2.3.
Figure 6.13 illustrates the improvement
that the tool provides. As of 2006-Dec, data for the two BI chips can also be
corrected in the same way, including a correction for serial CTI for the BI chips,
though the effects are more subtle. The resulting response
is very nearly uniform across the BI chips once this correction is made.

Figure 6.13: We plot energy resolution (FWHM in eV) versus
row number (CHIPY) for several cases. At the top, we show
Aluminum Kα (1.49 keV), while the bottom plot is for
Manganese Kα (5.9 keV). In each panel, the data points
represent the response of the I3 chip (upper curve, without CTI
correction; lower curve, with CTI correction), while the
lines represent the response of S3. Data were taken from
2009-May through 2009-Jul, on I3 node 3, and S3 node 0
(where the aimpoints are).

The ACIS energy resolution (here taken to mean the full width at half
maximum [FWHM] of a narrow spectral line) varies roughly
as the square root of the energy, and increases with distance from
the read-out. On FI chips (the I array and all the S array
chips except S1 and S3), the increase with CHIPY (row number)
is dramatic, as can be seen in Figure 6.13.
The spatial dependence on BI chips (S1 and S3) is much weaker,
but depends on both coordinates. To illustrate, we have measured the
resolution at the aimpoints and default offsets in effect 2009
(see Table 6.4),
using the aluminum (1.49 keV) and manganese (5.9 keV) Kα lines in
the external calibration source. Data were summed over three months
in the summer of 2009, and over an area of 32×32 pixels.
Note that these values are intended to be illustrative;
the spectral resolutions at the aimpoints
will differ from those in Table 6.4 because the
2009 aimpoints plus default offsets differ slightly from the default
aimpoints established for Cycle 18 and beyond; see
Section 4.5 for more detail.

Hot pixels and columns are defined to be those which produce a
spuriously high or saturated pulse-height for a large number of consecutive
frames of data. These depend on operating conditions such as
temperature. One should always refer to the CXC web site for the most
recent list.

Cosmic ray hits sometimes deposit so much charge that they appear
in the same pixels for a number of successive (or nearly successive) frames.
These cosmic ray afterglows are essentially temporary hot pixels, and
are removed by the hot pixel logic if they contain more than
approximately 8 events.

Aimpoints are the nominal positions on ACIS where the flux from a
point source with no target offset or SIM-Z translation is placed. There are
two aimpoints, indicated in Figure 6.1 - one
on the corner of I3 on the ACIS-I array (the ACIS-I aimpoint), and one
near the boundary between nodes 0 and 1 on S3 of the ACIS-S array (the
ACIS-S aimpoint).
Because of variations in the thermal conditions in the aspect subsystem
(see Section 5.4.3), the
actual position on the detector of a source with no offsets or
SIM-Z translations can vary somewhat from the default location.
Note also spacecraft dither will move the position of the
source on the detector in a Lissajous pattern centered on the
source position on the detector (see Section 6.11).

As of Cycle 18, permanent default aimpoints are specified for each
detector. Note that the aimpoint is not the same as
the optical axis, which is defined as the position of the
narrowest PSF. The chip coordinates of the aimpoints
and the current estimate for the optical axis position are listed
in Table 4.3. The aimpoints
are close to the optical axis, and have been adjusted to reduce the risk
of dither moving a source at the default aimpoint across a chip
edge (ACIS-I3) or across a node boundary (ACIS-S3) with the
expected aiming uncertainties.
Figures 4.26 and 4.27
show the default aimpoints, aimpoint error box, and aimpoint error box
plus dither. Further information on the absolute pointing and
aimpoint stability and uncertainties can be found in
Sections 5.4.3 and 4.5.

Approximate contours of constant encircled energy for ACIS-I and
ACIS-S observations at the aimpoints are shown in
Figures 6.14 and 6.15.

When ACIS is used with one of the transmission gratings, HETG, or LETG,
a SIM translation may be applied, and offsets may be applied
to the default aimpoints used for non-grating observations.
For further details,
see Chapter 8 (HETG), or Chapter 9 (LETG).

Figure 6.14: Approximate contours of constant 50%
encircled energy at 1.49 keV when the
ACIS-I
aimpoint is
selected. The dotted line is 1 arcsec, the dashed line is 1.5
arcsec. The remainder are at 1 arcsec intervals. The thicker solid lines
highlight the 5, 10, and 15 arcsec contours.

Figure 6.15: Approximate contours of constant 50%
encircled energy at 1.49 keV when the ACIS-S aimpoint is
selected. The dotted line is 1 arcsec, the dashed line is 1.5
arcsec. The remainder are at 1 arcsec intervals. The thicker solid lines
highlight the 5, 10, 15 and 20 arcsec contours.

Lastly, it should be kept in mind that for ACIS observations the
observatory is typically
dithered with a 16′′ peak-to-peak amplitude in the
Y and Z directions (see Section 6.11).

Unless specially requested, the spacecraft is dithered in a Lissajous
pattern during all observations. For observations with ACIS, the
default dither pattern spans 16 arcsec peak to
peak in both the Y and Z directions. The dither serves two
purposes: (1) to provide some exposure in
the gaps between the CCDs, and (2) to smooth out pixel-to-pixel
variations in the response. The effect of dither is removed during high-level
ground processing of the data. The exposure time in the gaps between
chips (and at the outside edges) will be less than that for the
remainder of the field. Default dither parameters are listed in
Table 5.4.

The approximate sizes of the various gaps between chips are shown in
Figure 6.1. Note that due to the way the
ACIS-I array CCDs are tilted to follow approximately the HRMA
focal plane curvature, the Y-gaps vary slightly with Z and the
Z-gaps vary slightly with Y.

ACIS has two operating modes: the Timed Exposure (TE) Mode,
and the Continuous Clocking (CC) Mode. One must select one or the
other for an observation as it is not possible to simultaneously
operate individual CCDs in different modes during a single observation.

A timed exposure refers to the mode of operation wherein a CCD collects
data (integrates) for a preselected amount of time - the
Frame Time. Once this time interval has passed, the charge from the
1024×1024 active region is quickly ( ∼ 41 ms) transferred to
the frame store region and subsequently read out through (nominally)
1024 serial registers.

Frame Times - Full Frames

Frame times are selectable within a range of values spanning the time
interval from 0.2 to 10.0 sec. If the data from the entire
CCD are utilized (full frame) then the nominal (and optimal!)
frame time, Topt, is 3.2 s. Selecting a frame time
shorter than the
nominal value (e.g. to decrease the probability of pile-up -
Section 6.15) has the consequence that there will be a
time during which no data are taken, "deadtime", as 3.2 s are required for
the frame store read-out process regardless of the frame time. The fraction of time
during which data are taken is simply the ratio of the selected frame
time to the sum of the nominal frame time and the 41 msec transfer time -
e.g. for a
new frame time of t ( < 3.2) secs, the fraction of time during which
data are taken is t/(Topt+0.041). We note, strictly speaking, the
full-frame time depends on how many CCDs are on
- see the equation below - but the differences are very
small. Finally, we note that selecting a frame time longer than
the most efficient value increases the probability of pile-up occurring
and is not recommended. In addition, with the standard ACIS dither,
a frame time longer than 3.2 s will add a blur to the PSF.

Frame Times & Subarrays

It is also possible for one to select a subarray - a restricted
region of the CCD in which data will be taken. A subarray is fully
determined by specifying the number of rows separating the subarray
from the frame store region (q) and the number of rows in the subarray
(n). Examples of subarrays are shown in Figure 6.16. The
nominal frame time for a subarray depends on (q), (n), and the total
number of CCDs that are activated (m) - see Table 6.5.
The nominal frame time is given by:

T(msec) = 41.12×m + 0.040 ×(m×q) + 2.85×n − 32.99.

(6.1)

As with full frames, selecting a frame time
less than the most efficient value results in loss of observing efficiency.
Frame times are rounded up to the nearest 0.1 sec, and can range from 0.2 to
10.0 sec.

When operating with only one chip, subarrays as small as 100 rows are
allowed (this permits 0.3 sec frame times which pay no penalty
in dead time). For multichip observations, the smallest allowed number of
rows is 128. For small subarrays, the aiming uncertainties plus dither
should be taken into account; see Section 6.10 and
references therein.

Table 6.5: CCD Frame Time (sec) for Standard Subarrays

Subarray

ACIS-I (no. of chips)

ACIS-S (no. of chips)

1

6

1

6

1

3.0

3.2

3.0

3.2

1/2

1.5

1.8

1.5

1.8

1/4

0.8

1.2

0.8

1.1

1/8

0.5

0.8

0.4

0.7

Figure 6.16: Examples of various subarrays. The heavy
dot in the lower left indicates the origin.

Trailed Images

It takes 40 μsec to transfer the charge from one row to another
during the process of moving the charge from the active region to the
frame store region. This has the interesting consequence that each
CCD pixel is exposed not only to the region of the sky at which the
observatory was pointing during the long (frame time) integration, but
also, for 40 μsec each, to every other region in the sky along
the column in which the pixel in question resides.
Figure 6.17 is an example where there are bright
features present so intense that the core of the PSF is suppressed
because of pile-up (see Section 6.15),
allowing the tiny contribution of the flux
due to trailing to be stronger than the (piled-up) core of the direct
exposure - hence the
trailed image is clearly visible. Trailed images are also referred to
as "read out artifacts", "transfer smear", or "out-of-time images".
The user needs to be aware of this phenomenon as it has implications
for the data analysis, including estimates of the background.
In some cases, the
trailed image can be used to obtain an unpiled spectrum and can also
be used to perform 40 microsec timing analysis of (extremely bright)
sources.

Figure 6.17: Trailed image of a strong X-ray
source. The core of the image is faint due to pile-up. Most events
here are rejected because of bad grades. The read-out direction is
parallel to the trail.

In some instances, it is desirable to have both long and short frame times.
If the exposure time is made very short, pile-up may be
reduced, but the efficiency of the observation is greatly reduced by the
need to wait for the full 3.2 sec (if six chips are clocked) for the
frame-store array processing.

With alternating exposure times, all CCDs are clocked in unison, but have
two exposure times. One (typically short) primary exposure of length
0.2 < tp < 10.0 sec is followed by k
standard exposures ts (for example, 3.2
sec if six chips are clocked).
Permissible values of k range from 1 to 15.
The short exposures are used
to reduce photon pile-up, and the long exposures are useful for studying the fainter
objects in the field of view. For example, a typical choice of long and short exposure times might
be 3.2 and 0.3 sec. If k = 3, ACIS would perform one 0.3 sec
exposure, followed by three exposures of 3.2 sec, repeating until the
total observing time expires.

If the duty cycle of long exposures is 1:k (short:long), the observing
efficiency η is then

The continuous clocking mode is provided to allow 3 msec timing at the
expense of one dimension of spatial resolution. In this mode, one
obtains 1 pixel × 1024 pixel images, each with an integration time of
2.85 msec. Details as to the spatial distribution in the columns are
lost - other than that the event originated in the sky along the line
determined by the length of the column.

In the continuous clocking mode, data are continuously clocked through
the CCD and frame store. The instrument software accumulates data into
a buffer until a virtual detector of size 1024 columns by 512 rows is
filled. The event finding algorithm is applied to the data in this
virtual detector and 3×3 event islands are located and recorded to
telemetry in the usual manner (Section 6.14.1). This procedure has the advantage that
the event islands are functionally equivalent to data accumulated in
TE mode, hence differences in the calibration are minimal. The row
coordinate (called CHIPY in the FITS file) maps into time in that a
new row is read from the frame store to the buffer every 2.85 msec.
This does have some minor impacts on the data. For example, since the
event-finding algorithm is looking for a local maximum centered in a
3×3 pixel area, it cannot find
events on the edges of the virtual detector. Hence, CHIPX cannot be 1
or 1024, and CHIPY cannot be 1 or 512. In
other words, events cannot occur at certain times separated by
512*2.85 msec or 1.4592 sec. Likewise, it is impossible for two
events to occur in the same column in adjacent time bins.

Event files for continuous-clocking mode data contain two
time columns, TIME, and TIME_RO. The TIME_RO records
the read-out times (effectively the time the virtual frame
is processed to identify events). The TIME column is
an estimate of the photon time of arrival at the detector,
based on the read-out time, the aspect solution, and the
specified sky location of the source. See the
acis_process_events documentation for information
on the generation of these columns. In order to ensure
that the TIMEs are as accurate as possible, it is best to
specify the source location to better than 0.5′′.
The read-out time can differ from the arrival time
by about 2.9 to 5.8 s, depending on the nominal location of the
source on the CCD and the dither of the spacecraft.

In general, the CCD bias, the amplitude of the charge in each pixel in
the absence of external radiation, is determined at every change of instrumental parameters or setup when ACIS is in place at the focus of the
telescope. These bias maps have proven to be remarkably stable and are
automatically applied in routine data processing.

The bias maps for continuous-clocking mode observations can be
corrupted by cosmic rays. If a cosmic ray deposits a lot of charge in
most of the pixels in one or more adjacent columns, the bias values
assigned to these columns will be too large. As a result, some
low-energy events that would have been telemetered will not be
telemetered because they do not satisfy the minimum pulse height
criterion and the spectrum of a source in the affected columns will be
skewed to lower energies. The BI CCDs are relatively insensitive to
the problem. A bias algorithm was
implemented in Cycle 6 to mitigate the problem.

Occasionally a cosmic ray produces an artifact in a bias map. The pipelines
search for these artifacts, and, when found, replace the bias map with another
from the same epoch. Work is in progress to use long-term average bias maps,
either when there are artifacts in the observation-specific bias map, or
for all observations.

In the first step in detecting X-ray events, the on-board processing
examines every pixel in the full bias-subtracted CCD image (even in
the continuous clocking mode (Section 6.12.3)) and selects
as events those with values that both exceed the event threshold and
are greater than all of their touching or neighboring pixels
(i.e., a local maximum in a 3×3 pixel detection island).
The pixel pattern, and thus the event grade,
is determined by which of the outer pixels in a 3×3 grid
centered on the initial pixel have values above the split-event
threshold. Depending on the grade, the data are then included in
the telemetry. On-board suppression of certain grades is used
to limit the telemetry bandwidth devoted to background events (see
Section 6.16.1).

The grade of an event is thus a code that identifies which pixels
within the 3×3 pixel island, centered on the local
charge maximum, are above certain amplitude thresholds. The thresholds
are listed in Table 6.2. Note that the local maximum
threshold differs for the FI and the BI CCDs. A "Rosetta Stone" to
help one understand the ACIS grade assignments is shown in
Figure 6.18, and the relationship to the ASCA grading
scheme is given in Table 6.6.

Figure 6.18: Schematic for determining
the grade of an event. The grade is determined by summing the numbers
for those pixels that are above their thresholds. For example, an
event that caused all pixels to exceed their threshold is grade 255. A
single pixel event is grade 0.

It is important to understand that most, if not all, calibrations of
ACIS are based on a specific subset of ACIS grades. This
standard
set comprises ASCA grades 0,2,3,4, and 6 - G(02346). In the absence
of pile-up, this particular grade selection appears to optimize the
signal-to-background ratio, but this conclusion depends on the
detailed spectral properties of the source. Further, most of the
scientifically important characteristics of ACIS (effective area,
sensitivity, point spread function, energy resolution, etc.) are
grade- and energy-dependent.

There are a number of telemetry formats available. Specifying a format
determines the type of information that is included in the telemetry
stream. The number of bits per event depends on which mode and which
format is selected. The number of bits per event, in turn, determines
the event rate at which the telemetry will saturate and data will be
lost until the on-board buffer empties. The formats available depend
on which mode (Timed Exposure or Continuous Clocking) is used. The
modes, associated formats, and approximate event rates at which the
telemetry saturates and one begins to limit the return of data, are
listed in Table 6.6. The formats are described in the
following paragraphs. Event "arrival time" is given relative to the
beginning of the exposure.

Table 6.7: Telemetry Saturation Limits

Mode

Format

Bits/event

Events/sec*

Number of Events

in full buffer

CC

Graded

58

375.0

128,000

CC

Faint

128

170.2

58,099

TE

Graded

58

375.0

128,000

TE

Faint

128

170.2

58,099

TE

Very Faint

320

68.8

23,273

*(includes a 10% overhead for housekeeping data)

Faint

Faint format provides the event position in detector coordinates, an
arrival time, an event amplitude, and the amplitude of the signal in each pixel in the
3×3 event island
that determines the event grade. The bias map is telemetered
separately. Note that certain grades may be not be included in the
data stream (Section 6.16.1).

Graded

Graded format provides event position in detector coordinates, an
event amplitude, the arrival time, and the event grade. Note that
certain grades may be not be included in the data stream
(Section 6.16.1).

Very Faint

Very Faint format provides the event position in detector coordinates,
the event amplitude, an arrival time, and the pixel values in a 5×5
island. As noted in Table 6.6, this format is only
available with the Timed Exposure mode. Events are still graded by the
contents of the central 3×3 island. Note that certain grades may be
not be included in the data stream (Section 6.16.1). This
format
offers the advantage of reduced background after ground processing
(see Section 6.16.2) but only for sources with low counting
rates that avoid both telemetry saturation and pulse pile-up.

Studies
(see
http://cxc.harvard.edu/cal/Acis/Cal_prods/vfbkgrnd)
of the ACIS background have shown that for
weak or extended sources, a significant reduction of background at low and
high energies may be made by using the information from 5×5 pixel
islands, i.e. very faint mode, instead of the faint mode 3×3
island. This screening results in a 1-2% loss of good events. CIAO 2.2 and later provides a tool to utilize the VF mode for screening
background events.
Please note that the RMF generation is the same for very faint mode
as it is for faint mode. See the "Why Topic"
http://cxc.harvard.edu/ciao/why/aciscleanvf.html.
The very faint mode screening also reduces the nonuniformity
of the non-X-ray background (see Section 6.16.1).

It is important to realize that VF mode uses more telemetry; the limit
is ∼ 68 cts/s, which includes the target flux and the full background
from all chips. Check the calibration web page
(http://cxc.harvard.edu/cal/Acis/Cal_prods/bkgrnd/current/background.html)
for a discussion of
background flares and the telemetry limit. In particular, review
Section 1.3 of the memo "General discussion of the quiescent and
flare components of the ACIS background" by M. Markevitch.

To reduce the total background rate and the likelihood of
telemetry saturation, VF observations should consider using no more
than 5 CCDs and an energy filter with a 12 keV upper cutoff.
Starting with Cycle 11, the default upper energy
cutoff has been decreased from 15 keV to 13 keV,
but VF mode observers should consider reducing this threshold
further to 12 keV.
If the target is brighter
than 5-10 cts/s, one has to take more drastic steps, such as turning off
more chips or using Faint mode.
Note that the CXC now encourages selection of 4 or fewer
required CCDs if the science can be accomplished with
4 or fewer CCDs, and at most 5 CCDs can be selected as required.
See Section 6.21.1
for more information on selection of required and optional CCDs.

This mode should not be used for observing bright sources (see the
discussion at the end of Section 6.16.1 for more detail).

Pile-Up results when two or more photons are detected as a single
event. The fundamental impacts of pile-up are: (1) a distortion of the
energy spectrum - the apparent energy is approximately the sum of two
(or more) energies; and (2) an underestimate as to the correct
counting rate - two or more events are counted as one. A simple
illustration of the effects of pile-up is given in Figure
6.19. There are other, somewhat more subtle, impacts
discussed below (Section 6.15.1).

The degree to which a source will be piled can be roughly estimated
using PIMMS. Somewhat more quantitative estimates can be obtained
using the pile-up models in XSPEC, Sherpa and ISIS . If the
resulting degree of pile-up appears to be unacceptable given the
objectives, then the proposer should employ some form of pile-up
mitigation (Section 6.15.3) as part of the observing
strategy. In general, pile-up should not be a problem in the
observation of extended objects, the Crab Nebula being a notable
exception, unless the source has bright knots or filaments.

Figure 6.19: The effects of pile-up at 1.49
keV (Al Kα) as a function of source intensity. Data were taken
during HRMA-ACIS system level calibration at the XRCF. Single-photon
events are concentrated near the pulse height corresponding to the Al
Kα line ( ∼ 380 ADU), and events with 2 or more photons
appear at integral multiples of the line energy.

There are other consequences of pile-up in addition to the two
principal features of spurious spectral hardening and underestimating
the true counting rate by under counting multiple events. These
additional effects are grade migration and pulse saturation, both of
which can cause distortion of the apparent PSF.

Grade migration

Possibly the most troubling effect of pile-up is that the grade
distribution for X-ray events may change. The change of
grade introduced by pile-up has become to be referred to as "grade
migration". Table 6.7 shows an example of grade
migration due to pile-up as the incident flux is increased. In this
simple test, which involved only mono-energetic photons, the largest
effect is the depletion of G0 events and the increase of G7 events.
In general, as the incident flux rate increases, the fraction of the
total number of events occupying a particular event grade changes as
photon-induced charge clouds merge and the resulting detected events
"migrate" to other grades which are not at all necessarily included
in the standard (G02346) set. If one applies the standard
calibration to such data, the true flux will be under-estimated.

Table 6.8: ASCA Grade Distributions for different incident fluxes at 1.49
keV (Al Kα, based on data taken at the XRCF during ground
calibration using chip I3)

Incident

Flux*

G0

G1

G2

G3

G4

G5

G6

G7

9

0.710

0.022

0.122

0.053

0.026

0.009

0.024

0.035

30

0.581

0.057

0.132

0.045

0.043

0.039

0.029

0.073

98

0.416

0.097

0.127

0.052

0.050

0.085

0.064

0.108

184

0.333

0.091

0.105

0.040

0.032

0.099

0.077

0.224

*arbitrary units

Pulse Saturation

One consequence of severe instances of pile-up is the creation of a
region with no events! In this case, the pile-up is severe enough that
the total amplitude of the event is larger than the on-board threshold
(typically 13 keV) and is rejected. Holes in the image can also be
created by grade migration of events into ACIS grades (e.g. 255) that
are filtered on-board.

PSF distortion

Obviously the effects of pile-up are most severe when the flux is highly
concentrated on the detector. Thus, the core of the PSF suffers more
from pile-up-induced effects than the wings. Figure 6.20
illustrates this point.
Because the core is suppressed, the PSF profile
appears less peaked and (apparently) broader than would be the case
if pile-up were negligible.

It is clearly important in preparing a Chandra observing proposal to
determine if the observation will be impacted by pile-up, and if so,
to decide what to do about it (or convince the peer review why the
specific objective can be accomplished without doing anything). There
are two approaches to estimating the impact of pile-up on the
investigation. The most sophisticated uses the pile-up models in
XSPEC, Sherpa, and ISIS to create a simulated data set which can be
analyzed in the same way as real data. A less sophisticated,
but very useful, approach is to use the web version of PIMMS to
estimate pile-up, or to use Figures 6.20 and
6.21.

Figure 6.20: The effects of pile-up on the radial
distribution of the PSF are illustrated. These data were
taken during ground calibration at the XRCF. The specific "OBSIDs",
the counting rate per CCD frame ("c/f"), and the "pile-up fraction"
as defined in Section 6.15.2 are given in the inset.

Figure 6.21: The pile-up fraction as a
function of the the counting rate (in the absence of pile-up in units
of photons/frame). The solid line is for on-orbit, the dashed line and
the data points are for, and from, ground-based data respectively. The
difference between the ground and flight functions are a consequence
of the improved PSF on-orbit, where gravitational effects are
negligible - see Chapter 4. Note that when pile-up occurs
there are two or more photons for each event, so the fraction of
events with pile-up is always less than the fraction of photons with
pile-up.

Figure 6.22: MARX simulations of the effect of pile-up on the shape of the spectrum.
The true (solid line) and the detected (dotted line) spectra are shown for
four different viewing angles. The corresponding "pile-up fractions"
- see Section 6.15.2 - are 46%, 40%, 15%, and 2% as
the image is moved progressively further off-axis.

Simple Pile-Up Estimates
The pile-up fraction is the ratio of the number of detected events that
consist of more than one photon to the total number of detected
events. An estimate of the pile-up fraction can be determined from
Figure 6.21. The algorithm parametrizes the
HRMA+ACIS PSF in terms of the fraction of encircled energy that
falls within the central 3×3 pixel event detection cell, and assumes that the remaining energy is uniformly
distributed among the 8 surrounding 3×3 pixel detection
cells. The probabilities of single- and multiple-photon events are
calculated separately for the central and surrounding detection cells
and subsequently averaged (with appropriate weighting) to obtain the
pile-up fraction as a function of the true count rate - the solid
line in Figure 6.21. The model was tested against data
taken on the ground under controlled conditions - also shown in
Figure 6.21.

As a general guideline, if the estimated pile-up fraction is > 10%,
the proposed observation is very likely to be impacted. The first
panel (upper left) in Figure 6.22 qualitatively
illustrated the effect on a simulated astrophysical X-ray spectrum.
However, the degree of pile-up that is acceptable will depend
on the particular scientific goals of the
measurement, and there is no clear-cut tolerance level. If one's
scientific objective demands precise flux calibration, then the pile-up
fraction should probably be kept well below 10%.

The PIMMS tool provides the pile-up fraction using the algorithm
described here, both for direct observation with ACIS and also for
the zeroth-order image when a grating is inserted.

Simulating Pile-Up

John Davis developed an algorithm for modeling the effects
of pile-up on ACIS spectral data. The algorithm has been implemented
as of XSPEC V11.1 and Sherpa V2.2. The algorithm can be used to
attempt to recover the underlying spectrum from a source, or to
simulate the effects of pile-up for proposal purposes.

The algorithm has been tested by comparing CCD spectra with grating
spectra of the same sources. Care should be taken in applying the
algorithm - for example, using the appropriate regions for extracting
source photons and avoiding line-dominated sources. A description of
the algorithm can be found in Davis 2001 (Davis, J.E. 2001, ApJ, 562,
575). Details on using the algorithm in Sherpa are given in a
Sherpa "thread" as of CIAO V2.2 on the CXCCIAO web page:
http://cxc.harvard.edu/ciao/.

By cutting back on CCD exposure time,
the probability of pile-up decreases. The user is advised to select the
best combination of a subarray and frame time to avoid losing
data as discussed in Section 6.12.1.

• Use the Alternating Exposure option:

This option simply
alternates between exposures that are subject to pile-up and those that
are not. The capability was originally developed for use with certain
grating observations to allow one to spend some time obtaining useful
data from a zeroth order image, which would otherwise be piled up.
See Section 6.12.2.

• Use CC mode

If two-dimensional imaging is not required,
consider using CC mode (Section 6.12.3).

• Insert a transmission grating:

Inserting either
the HETG (Chapter 8) or the
LETG (Chapter 9) will
significantly decrease the counting rate as the efficiency is
lower. The counting rate in the zero order image may then be low
enough to avoid pile-up.

• Offset point:

Performing the observation with the source
off-axis spreads out the flux and thus decreases the probability of
pile-up at the price of a degraded image. Figure 6.22
illustrated the impact.

• Defocus:

The option is only listed for completeness, the option
is not recommended or encouraged.

There are three components to the on-orbit background. The first is
that due to the cosmic X-ray background (a significant fraction of
which resolves into discrete sources during an observation with
Chandra). The second component is commonly referred to as the charged
particle background. This arises both from charged particles,
photons, and other neutral particle interactions that ultimately deposit
energy in the instrument. The third component is the "read-out
artifact" which is a consequence of the "trailing" of the target
image during the CCD read-out; it is discussed in
Section 6.12.1.

The background rates differ between the BI and the FI chips, in
part because of differences in the efficiency for identifying charged
particle interactions. Figure 6.23 illustrates why.

Figure 6.23: Enlarged view of an area of a FI chip I3 (left) and
a BI chip (right) after being struck by a charged particle. There is far
more "blooming" in the FI image since the chip is thicker. The
overlaid 3×3 detection cells indicate that the particle impact on the
FI chip produced a number of events, most of which end up as
ASCA Grade 7, and are thus rejected with high efficiency. The equivalent
event in the BI chip, is much more difficult to distinguish from an
ordinary X-ray interaction, and hence the rejection efficiency is
lower.

Beginning in 2002-Sep
and continuing until 2012-Jun,
observations have been carried
out with the ACIS in the stowed position, shielded from the sky by
the SIM structure, and collecting data in normal imaging TE VF mode at
-120C. Chips I0, I2, I3, S1, S2,
S3 were exposed. The SIM position was chosen so that the on-board calibration
source did not illuminate the ACIS chips. This allowed us to
characterize the non-celestial contribution to X-ray background (i.e.,
from charged particles). The resulting spectra from different chips
are shown in Figure 6.24. Chip S2 is similar
to the ACIS-I chips (denoted I023 in the figure) and not
shown for clarity.

In addition, in July-September 2001, Chandra performed several short
observations of the dark Moon, which blocks the cosmic X-ray background.
The dark Moon and stowed background spectra were indistinguishable
(except for short periods of flares and variable Oxygen line emission
in the Moon observations). We have not observed any background flares
in the stowed position. Thus, the ACIS-stowed background is a good
representation of the quiescent non-X-ray background in the normal
focal position and can be used for science observations.

Figure 6.24: Energy spectra of the charged
particle ACIS background with ACIS in the stowed position (a 50 ks
exposure taken in 2002-Sep; standard grade filtering, no VF
filtering). Line features are due to particle-induced fluorescence of
material in and surrounding the focal plane.

The flight-grade distributions in early measurements of the non-X-ray
background for the two types of CCDs are shown in
Figure 6.25. Although subsequent to these early
measurements the CCD temperature has been lowered and the
FI devices suffered the effects of the radiation damage, the
background is still dominated by the same grades. Based on these data,
events from flight grades 24, 66, 107, 214, and 255 are routinely
discarded on-board. The total rate of the discarded events is
available in the data stream. The remaining non-X-ray events
telemetered to the ground are still dominated (70-95%) by other bad
grades. They are not discarded on-board because some of them may turn
out to be valid X-ray events after ground processing.

For data taken using the Very Faint (VF) telemetry format
(Section 6.14.2), the non-X-ray background can be reduced
in data processing by screening out events with significant flux in
border pixels of the 5×5 event islands. This screening leaves the data
from faint sources essentially the same while reducing
the FI background at different energies: a factor of
1.4 (E > 6 keV);
1.1 (1−5 keV); and
2 (near ∼ 0.5 keV).
For the BI chips the reductions are:
1.25 (E > 6 keV);
1.1 (1−5 keV); and
3 (near 0.3 keV).
In addition, the spatial nonuniformity of the non-X-ray background may
be reduced by VF screening; see the background uniformity memo at
http://cxc.harvard.edu/cal/Acis/detailed_info.html#background.
The screening
algorithm has been incorporated into the CIAO tool
acis_process_events. Further discussion may be found at
http://cxc.harvard.edu/cal/Acis/Cal_prods/vfbkgrnd/index.html.

Proposers should be aware that telemetry saturation occurs at lower
count rates for observations using the VF format, so they may need to
take steps to limit the total ACIS count rate (see
Section 6.16.2). Proposers should also be aware that if there
are bright point sources in the field of view, the screening criterion
discussed above is more likely to remove source events due to pile-up
of the 5×5 pixel event islands. Point sources should have count rates
significantly less than 0.1 cts/sec to be unaffected. However, there
is no intrinsic increase of pile-up in VF data compared to Faint mode,
and the screening software can be selectively applied to regions,
excluding bright point-like sources. Thus the use of VF mode is
encouraged whenever possible.

Figure 6.25: Fraction of ACIS background events as a function of grade from
early in-flight data for an FI chip (S2) (top) and a BI chip (S3)
(bottom).

In real observations, two more components to the background come into
play. The first is the cosmic X-ray background which, for moderately
long ( ∼ 100 ks) observations will be mostly resolved into discrete
sources (except for the diffuse component below 1 keV) but,
nevertheless, contributes to the overall counting rate. The second is
a time-variable "flare" component caused by any charged particles
that may forward-scatter from the HRMA mirrors and have sufficient
momentum so as not to be diverted from the focal plane by the
magnets included in the observatory for that purpose, or
from secondary particles (Section 6.16.3).
Figure 6.26 compares flare-free
ACIS-S3 spectra of the non-X-ray (dark Moon) background and a
relatively deep pointing to a typical region of the sky away from
bright Galactic features.

Figure 6.26: ACIS-S3
spectrum of the non-X-ray background (large
crosses) overlaid on the quiescent blank sky spectrum. Small
crosses show the total sky spectrum, while squares show the
diffuse component left after the exclusion of all point
sources detectable in this 90 ks exposure. Standard grade
filtering and VF filtering are applied. The background and
blank-sky spectra are normalized to the same flux in the 10-12
keV band.

Estimates for the mid-2015 quiescent detector+sky background
counting rates in various energy bands and for the standard
good grades are given in Table 6.8. Insertion of the
gratings makes little measurable difference in the
background rates, but it does block the background
flares. The lower-energy rates are very approximate as they
vary across the sky. The rates are slowly changing on the
timescale of months, so Table 6.8 can only be used for rough
sensitivity estimates. Table 6.9 gives total background count rates for
each type of chip, including all grades that are telemetered (see
Section 6.14.1 and 6.16.1),
and can aid in estimating the probability of telemetry saturation.

ACIS Background rates (cts/s/chip)

Energy

Band (keV)

I0

I1

I2

I3

S1

S2

S3

S4

0.3-10.0

0.31

0.32

0.31

0.32

1.62

0.30

0.80

0.33

0.5-2.0

0.05

0.06

0.06

0.06

0.14

0.06

0.11

0.08

0.5-7.0

0.17

0.18

0.17

0.18

0.43

0.17

0.30

0.20

5.0-10.0

0.17

0.18

0.17

0.18

1.18

0.17

0.50

0.16

10.0-12.0

0.10

0.10

0.10

0.10

0.94

0.10

0.63

0.10

Table 6.8: Approximate on-orbit standard grade background counting
rates (2015, August to mid-October). The rates are cts/s/chip,
using only ASCA grades 02346, no VF filtering,
excluding background flares and bad
pixels/columns and celestial sources identifiable by eye.
These values can be used for sensitivity calculations

Period

Aug 1999

2000-2003

2009

2015

Upper E cutoff

15 keV

15 keV

15 keV

15f keV

13 keV

12 keV

10 keV

Chip S2 (FI)

10

6.3

10.7

6.5

5.3

5.2

4.8

Chip S3 (BI)

11

7.7

14.7

10.2

6.9

6.3

4.5

f Scaled from 13 keV rate.

Table 6.10: Typical total quiescent
background rates (cts/s/chip; 2015, August to mid-October),
including all grades that are telemetered (not
just standard ASCA grades), by chip type and upper energy
cutoff. These values can be used to estimate the probability of
telemetry saturation.

The background rates declined in 1999-2000,
stayed flat through end of 2003, increased until the end of 2009,
dropped thereafter, and finally flattened more recently, approximately
following the (inverse of) the solar cycle (with an offset of about a year);
see Figure 6.27. By late 2015, the ACIS-S2 rates
remained at about the level of the low rates of 2001-2003.
The ACIS-S3 rates had also leveled off, but at somewhat higher rates
than in 2001-2003. The trend indicates that both rates are beginning
to climb, and in 2016 the ACIS-S2 and ACIS-S3 rates may rise further.
The rates are very approximately anticorrelated with the solar
cycle, and the rates are expected to increase as the sun becomes
less active.
The proposer should be cautious in making assumptions about
the future background rates for planning purposes.

Figure 6.27: Total telemetered background rates (including all grades
and the upper event cutoff at 15 keV) for chips S2
(FI) and S3 (BI) as a function of time. Vertical dashed lines
are year boundaries.

For aid in data
analysis and planning background-critical observations, the
CXC has combined a number of deep, source-free, flare-free
exposures (including all components of the background) into
background event files for different time periods. These
blank-sky datasets, along with the detector-only
(ACIS-stowed) background files (Section 6.16.1), can be
found in the CALDB.

In general, the background counting rates are stable during an
observation. Furthermore, the spectral shape of the non-X-ray
background has been remarkably constant during 2000-2015 for FI
and, to a lower accuracy, for BI chips, even though the overall background rate showed secular
changes by a factor of 1.5. (For chip S3, the shape has been constant during 2000-2005,
but a small change has been observed since late 2005.) When the quiescent background spectra
from different observations are normalized to the same rate in the
10-12 keV interval, they match each other to within ±3% across the
whole Chandra energy band. The previous discussion assumes that the upper
threshold is set to 13 keV.

Occasionally, however, there are large variations (flares), as
illustrated in Figure 6.28.
Figure 6.29
shows the frequency of such variations when compared to the quiescent
background.
An average fraction of the exposure affected by flares
above the filtering threshold used for the blank-sky
background datasets (a factor of 1.2 above the nominal rate)
was about 6% for FI chips and up to 1/3 for BI chips during
the first few years of the mission. The average fraction of exposure
affected by flaring has
declined with time, and was practically zero for a long stretch.
Recently, the flare frequency has been increasing, but at present
(late 2015) it has not reached the frequency seen early in the mission.
Thus, given that the quiescent background
in FI chips is also lower than that in S3, background-critical
observations may best be done with ACIS-I.

Several types of flares have been identified, including flares that
are seen only in the BI chips, and flares that are seen in both
the FI and BI chips. Figure 6.30 shows
the spectra of two of the
most common flare species. Both flares have spectra significantly
different from the quiescent background. In addition, the
BI flares have spatial distribution very different from that of the
quiescent background. The BI flares produce the same spectra in S1
and S3.

Users should note that the total counting rate can significantly
increase during a flare (although flare events are almost exclusively
good-grade so the total rate does not increase by as large a factor as
the good-grade rate; details can be found at
http://cxc.harvard.edu/cal/Acis/Cal_prods/bkgrnd/current). If
the probability of telemetry saturation is significant, users of
ACIS-I might consider turning off the S3 chip. However, if ACIS-S is
used in imaging mode,
and flaring is a concern, the CXC recommends
that both BI chips be
turned on. The advantage is that for most types of flares S1 can be
used to determine the flare spectrum which can then be subtracted from
the spectrum obtained with S3.

Figure 6.28: An example of the
ACIS background counting rate versus time - BI chip (S3; top curve) and
an FI chip (I2; bottom curve). These are for the standard grades and
the band from 0.3 - 10 keV.

Figure 6.29: An estimate of the
cumulative probability that the ratio of the background counting rate
to the quiescent background counting rate is larger than a given
value. Upper plot for a representative FI chip - S2, and the lower
curve for a representative BI chip - S3. The vertical dotted line is a
limiting factor 1.2 used in creating the background data sets.
These probabilities are relevant for the archival data from the first
2-3 years of the mission. After declining with time to almost zero for
a number of years, the amount of flaring has increased as we approach
the solar maximum, but there is less flaring than was seen in the first few
years of the mission.

Figure 6.30: Spectra of different background
flares in chip S3. Thick crosses show a common flare species that
affects only the BI chips. Thin crosses show one of the several less
common flare species that affect both the BI and FI chips. Note how
both these spectra are different from the quiescent spectrum (see
Figure 6.26).

Apart from compressing the data into one dimension
(Section 6.12.3), there is essentially no difference in the
total background in CC mode and that encountered in the timed exposure
mode. The background per-sky-pixel, however, will be 1024 times
larger, since the sky-pixel is now 1 × 1024 ACIS pixels.

The sensitivity for detecting objects is best estimated using the various
proposal tools such as PIMMS, MARX, etc. The "Chandra Proposal
Threads" web page gives detailed examples of how to use
these tools
(http://cxc.harvard.edu/proposer/threads).

Pre-Flight radiation tests have shown that ∼ 200 krads of X-ray
photon dose will damage the CCDs. The mechanism for the
damage is the trapped ionization in the dielectric silicon oxide and
nitride separating the gates from the depletion region. Since the
charge is trapped, the damage is cumulative. Because the structure of
the BI CCDs differs significantly from that of the FI CCDs, the two types of
chips have different photon dose limitations. Specifically, the BI CCDs are more than 25 times as tolerant to a dose of X-ray photons as
compared to the FI CCDs since the former have ∼ 40 μm of bulk Si
protecting the gate layer.

Simulations of astrophysical sources have yielded a very conservative,
spectrally-averaged, correspondence of 100 cts/pix = 1 rad. By
'counts' in this context we mean all photons that impinged on the
detector, whether or not they were piled-up or discarded on-board.

In consultation with the IPI team, the CXC has adopted the following
mission allowances, per pixel of the two types of chips:

If your observation calls for observing a bright point-like source
close to on-axis, we suggest you use the MARX simulator (with the
parameter DetIdeal=yes & dither, typically, on) to calculate whether
your observation may reach 1% of the above mission limits in any one
pixel. If so, please contact the CXC HelpDesk (http://cxc.harvard.edu/helpdesk/) to custom design
an observational strategy which may accommodate your science aims,
while maintaining the health & safety of the instrument.

Many of the spacecraft components have been reaching
higher temperatures over the course of the mission
because of changes in the insulating layers on the
exterior surfaces of the Chandra spacecraft. The
ACIS electronics and Focal Plane (FP) temperatures can
reach their operational limits depending on the orientation
of the spacecraft and the number of operating CCDs.
The number of operating CCDs and/or the duration of
observations is limited for observations with
solar pitch angles less than 60° and solar
pitch angles greater than 130°. Starting in
Cycle 15, the CXC has been encouraging GOs to
select 4 or fewer required CCDs unless the science
objectives require 5 or 6 CCDs. This is a departure
from the previous recommendation of the CXC which asked
GOs to select 6 CCDs to maximize the archival utility
of the observations.

Please note that the solar pitch
angles specified here, while they correspond to the
pitch range boundaries given in other CXC documents,
are practical representations of more complex physical
behavior and should be understood as approximate.

6.19.1 Background Information on the Thermal Environment and its Impact on ACIS

Several components on the Chandra spacecraft have
reached elevated temperatures at a variety of pitch angles.
Figure 6.31 displays approximate
pitch ranges and the components sensitive in those ranges.
Within ACIS, the Power Supply and Mechanism Controller (PSMC)
heats at pitch angles less than 60° and the
Focal Plane (FP), Detector Electronics Assembly (DEA),
and Digital Processing Assembly (DPA) heat at angles larger
than 130° (see Figure 6.32).
Under current conditions, and assuming an
initial PSMC temperature of less than +30 C, observations at
pitch angles less than about 60°, longer than 50 ks,
and with 6 CCDs operating are likely to approach or exceed the
PSMC thermal limit. Likewise observations at pitch angles larger
than about 130°, longer than about 50 ks, and with
6 CCDs operating are likely to approach or exceed the DPA and
DEA thermal limits (Figures
6.33 and
6.34).
Finally, the ACIS FP temperature will increase
above the desired operating temperature of -119.7 C for
observations with pitch angles larger than 130°. Each
of these temperatures can be reduced by reducing the number of
operating CCDs. The Operations team may turn off one or two
optional CCDs if a total of 5 or 6 CCDs are selected (where
"total" means the sum of required plus optional CCDs; see
section 6.21.1)
if the thermal models indicate that operational limits might be
reached for a given observation.

Figure 6.32 shows
the PSMC temperature as a function of spacecraft pitch angle using
data from a series of observations from between 2007 and 2014,
and for which the exposure times are at least 50 ks.
The figure illustrates the increase in PSMC temperature for
observations at low pitch angles.
Observations using 6 chips are plotted as small (red) stars,
and those using 5 chips are plotted as large (blue) circles;
variations in the maximum temperature
at a particular pitch angle within these two cases correspond
primarily to variations in the starting temperatures for the
observations. The maximum allowable PSMC temperature is
indicated by the light gray horizontal line (yellow line in the
online version). The PSMC temperature
has been controlled since 2008 through the use of a model
which predicts the temperature for a given week, given the
mix and timing of observations. If the
predicted temperatures exceed the planning limits, adjustments
are made such as turning off an optional CCD, splitting an
observation, or rescheduling an observation at a more favorable
pitch angle. It is evident from the plot that using one
less CCD can reduce the temperature by a few degrees, thus allowing
somewhat longer observations to be carried out at low pitch angles.

In tail-Sun orientations (pitch angles larger than 130°),
the ACIS FP temperature, the Detector Electronics Assembly (DEA)
temperature, and the Digital Processing Assembly (DPA) temperature
can warm outside of the desired range. Each of these three
temperatures can be reduced by reducing the number of CCDs
being clocked. Figure 6.33
displays the DPA temperature as a function of pitch angle for
5 and 6 CCD configurations, and
Figure 6.34
shows the corresponding data for the DEA.

The DPA temperature has been controlled
since 2012 through the use of a model which predicts the
temperature for a given week. If the predicted temperatures
exceed the planning limits, adjustments are made such as
turning off an optional CCD, splitting an observation,
or rescheduling an observation at a more favorable pitch angle.

Figure 6.31: Pitch sensitivity of spacecraft components
[See the online Proposers' Observatory Guide for a color version of this figure]

Figure 6.32: ACIS PSMC temperature as a function of spacecraft pitch angle
6-CCD observations are indicated by small (red) stars,
and 5-CCD observations are indicated by larger (blue) circles.
[See the online Proposers' Observatory Guide for a color version of this figure]

Figure 6.33: DPA temperature as a function of pitch angle for 5 and 6 CCD configurations.
This figure shows the variation of DPA temperature with pitch angle.
6-CCD observations are indicated by small (red) stars,
and 5-CCD observations are indicated by larger (blue) circles.
[See the online Proposers' Observatory Guide for a color version of this figure]

Figure 6.34: DEA temperature versus pitch angle for 5 and 6 CCD
observations.
This figure shows the variation of DEA temperature with pitch angle.
6-CCD observations are indicated by small (red) stars, and
5-CCD observations are indicated by larger (blue) circles.
[See the online Proposers' Observatory Guide for a color version of this figure]

Chandra has successfully observed several solar system objects, including
Venus, the Moon, Mars, Jupiter and several comets. Observations of planets
and other solar system objects are complicated because these objects move
across the celestial sphere during an observation and the optical light from
the source can produce a significant amount of charge on the detectors (this
is primarily an issue for ACIS-S observations). Some information regarding
observation planning and data processing is given here. Users are encouraged
to contact the CXC for more detailed help.

Chandra cannot observe the Sun for obvious reasons. Chandra has conducted
observations of the Moon earlier in the mission, but observations of the Moon
with ACIS are currently not allowed. The concern is that the
bright flux of optical and
UV photons could potentially polymerize the contaminant on the ACIS filters.
Observations of the dark portion of the Moon are not allowed since there is
a risk that the Sun-illuminated portion of the Moon might encroach upon the
FOV during the observation. For similar reasons, ACIS
observations of the Earth (including the dark portion) are not allowed.
See Section 3.3.2 in Chapter 3 and
Chapter 5 for further discussion on avoidances and
constraints.

Any solar system object other than the Sun, the Earth, the Moon,
and Mercury can be observed with ACIS-I, subject to the avoidances
discussed in Section 3.3.2 and Chapter 5.
Previous solar system
observations with ACIS-I have not
revealed significant contamination from optical light. However, proposers
are encouraged to work with the CXC when planning the specifics of a
given observation. Since the source moves across the celestial sphere in
time, an image of the event data will exhibit a "streak" associated with
the source. The CIAO tool sso_freeze can be used to
produce an event data file with the motion of the source removed.

Any solar system object other than the Sun, the Earth, the Moon,
and Mercury can be observed with ACIS-S, subject to the avoidances
discussed in Section 3.3.2 and Chapter 5.
The ACIS-S array can be used with or without a grating.
The BI CCDs are more
sensitive to soft X-rays than the ACIS-I array CCDs,
but the entire ACIS-S array
suffers from the disadvantage that its OBF is thinner than for ACIS-I and
may transmit a non-negligible flux of visible light onto the CCDs. It is
thus necessary to estimate the amount of charge produced in the CCDs due to
the optical light. More detailed information can be found at
http://cxc.harvard.edu/cal/Hrma/UvIrPSF.html
and from the CXC via HelpDesk
(http://cxc.harvard.edu/helpdesk/).

If the optical light leak is small enough, it can be mitigated by simply
shortening the frame time. This leads to a linear drop in the number of ADU
due to optical light. If possible, VF mode should be used, since in this
mode the outer 16 pixels of the 5×5 region allows a "local"
bias to be subtracted from the event to correct for any possible light
leakage. However,
see the warnings in Section 6.14.2.

The optical light also invalidates the bias taken at the beginning of the
observation if a bright planet is in the field. It is therefore desirable
to take a bias frame with the source out of the field of view. This bias
map is useful even when processing 5×5 pixels in VF mode since it can be
employed as a correction to the local average "bias" computed from the
16 outer pixels, thereby correcting for hot pixels, cosmetic defects etc.

A more sophisticated approach to dealing with excess charge due to optical
light is to make an adjustment to the event and split thresholds. Event
grades are described in more detail in Section 6.14.1. Excess charge (in
ADU) due to optical light will be added to the event and split counters
on-board. Without an adjustment to the thresholds (or a large enough
threshold), many of the X-ray events may have all nine pixels of a
3×3 pixel event detection cell above the split threshold, in which case
the event will not be telemetered to the ground. If the adjustment is too
large, X-ray events may not be detected because they may not exceed the
event threshold.

Users should be aware that if the detection thresholds are adjusted,
standard CXC processing of planetary data will give inaccurate estimates
of event pulse heights and grades. To analyze such data, a thorough
understanding of the energy calibration process and manual massaging of
the data will be required.

This section describes the various inputs that either must be, or can
be, specified to perform observations with ACIS. The
subsections are organized to match the RPS form. We have added some
discussion as to some of the implications of the possible choices. As
emphasized at the beginning of the Chapter, ACIS is moderately
complex and the specific characteristics of the CCDs and their
configuration in the instrument lead to a number of alternatives for
accomplishing a specific objective - detailed trade-offs are the
responsibility of the observer. For example, it might seem obvious that
observations of a faint point source may be best accomplished by
selecting the ACIS-S array with the aim point on S3, the BI device
that can be placed at the best focus of the telescope, and the
CCD with the best average energy resolution. On the other hand,
perhaps the science is better served by offset pointing (by a few
arcmin) the target onto S2, very near to the frame store, where the
FI energy resolution is better than that of S3. Or,
if the object is very faint - so that the number total number of
photons expected is just a handful (not enough to perform any
significant spectroscopy) - the advantage of S3 may not be so
obvious considering the smaller field of view and its higher
background rate, and perhaps the ACIS-I array, which would optimize
the angular resolution over a larger field, may be more attractive.

Some ACIS input parameters must be specified: the
number and identity of the CCDs to be used, the Exposure Mode, and
the Event Telemetry Format. If pile-up and telemetry saturation are not
considered to be a problem for the observation, then these are the
only parameters that need to be specified.

Number and Choice of CCDs
The RPS requires the observer to specify the desired aimpoint
and to identify the CCDs they want to use.
There has been a change in policy concerning the number of CCDs
that may be simultaneously operated. Prior to Cycle 13, we
encouraged the use of 6 CCDs to facilitate serendipitous
detections. Now, for thermal reasons, the CXC encourages
that you specify no more than
4 required CCDs
if your science objectives can
be met with 4 or fewer CCDs. If 4 CCDs are specified as required,
the GO may specify 1 or 2 optional CCDs to bring the total of
required plus optional CCDs to 5 or 6 CCDs, but the GO should be
aware that optional CCDs may be turned off once the thermal
environment is known at the time the STS is generated.
GOs may still request 5 CCDs as required, but they should only do
so if their science objectives require 5 CCDs. GOs may also still
request a total of 6 CCDs, but only 5 CCDs may be required and
one CCD must be specified as optional.

Using fewer than 6 CCDs is beneficial in keeping ACIS Power Supply
Mechanism Controller (PSMC) and the ACIS focal
plane (FP) and electronics temperatures within the required operating
ranges.
See Section 6.19
for further information on thermal limitations for ACIS
and the number of operating CCDs.

If the observer requires the optimal spectral response
and most accurate gain calibration (provided by a cold and stable focal
plane temperature), they may wish to select a total of 4 or fewer CCDs.
For example, if the observer is using the ACIS-I array for imaging,
they could select just the four I-array CCDs. If the observer is
using S3 for imaging, they could select just S2, S3, S4, and I3,
or they could select just S1, S2, S3, and S4.
How choices are designated in the RPS form is discussed in the
next subsection.

Choosing Optional CCDs & Optional CCD Policy
The observer may specify that a given CCD must be on for an
observation by entering "Y" for that CCD at the appropriate
place in the RPS form. If there are CCDs that the observer
does not require for their science, they should enter "N".
If there are CCDs that the observer would prefer to have
turned on should thermal conditions allow it, the rank-ordered designations
"Off1", "Off2", up to "Off5" should be used. The CCD designated
as "Off1" would be the first one to be turned off, and the
CCD designated as "Off5" would be the last that would be turned off.

Even if the science requires 6 CCDs, the observer should set the
designation for the 5 most useful as "Y" and the least useful
as "Off1". The proposer should also add a comment in the RPS form
that 6 CCDs are required for the science. If the proposal is
accepted, the observer may work with their User Uplink Support Scientist
to change "Off1" to 'Y'. The CXC will make its best
effort to try to schedule the observation under the appropriate
thermal conditions. Note that this scheduling is complicated,
might require breaking the observation into separate pointings,
and may be impossible to accomplish. The observer should discuss the
configuration with Uplink Support and, if there are difficulties
in assessing which CCDs should be optional, please contact the
CXC HelpDesk (http://cxc.harvard.edu/helpdesk/). Should it be possible to accommodate the observer's
request for 6 CCDs, the observation will most likely take place
at solar pitch angles between about 60° and 130°.

Some Recommended Chip Sets
Observers should specify the chip set that is best for their primary
science. The following suggestions have proven to be popular,
and would facilitate a more useful and homogeneous archive.

Given the current thermal performance of the spacecraft,
it is possible that an optional CCD would be turned off.

The rationale for the first ACIS-I imaging configuration
(Figure 6.35)
is that,
in the unlikely event of major background flares, telemetry
might saturate more rapidly if S3 were on and the focal-plane
temperature and electronics temperatures will be lower
with only 4 or 5 CCDs on. In addition, since S2 is further
from the ACIS-I aimpoint on I3, data on S2 may provide
nearby local (FI) background if the target is diffuse emission
which does not extend as far as S2.

For the second ACIS-I imaging configuration
(Figure 6.36),
the rationale is that S3 is generally more sensitive and closer
to the ACIS-I aimpoint, and so
more sensitive to serendipitous source detection.

The rationale for the third ACIS-I imaging configuration
(Figure 6.37)
is that it is desired to have both
S2 and S3 on, but it is not required.
Given the current thermal performance of the spacecraft,
it is probable that one of the two optional CCDs would be turned off and
possible that the second would also be turned off.
If thermal conditions require optional CCDs to be turned off,
the CCD specified "OFF1" would be the first to be turned off.

For deciding on chipsets for ACIS-S imaging, several factors come into
play. In general, chips closest to the S3 aimpoint would be selected
as required, while those farthest from the aimpoint (where the PSF
is degraded) would selected as optional, and be turned off first
if necessary.
Given the current thermal performance of the spacecraft,
it is possible that an optional CCD would be turned off.

The rationale for the first ACIS-S imaging configuration
(Figure 6.38)
is that S1 will have a higher count rate than an FI CCD in the event
of a background flare and thus it might be desirable to turn S1 off.

The rationale for the second ACIS-S imaging configuration
(Figure 6.39)
is that the observer may want to use S1 to model the background on
S3, and it is desirable to have S4 turned off. S4 also has
has significant noise streaks with resulting decreased sensitivity.
so some users may prefer to specify I2 as optional instead of S4.

The rational for the third ACIS-S imaging configuration
(Figure 6.40
is that it is desirable to have 6 CCDs on, but not required.
Given the current thermal performance of
the spacecraft, it is probable that the first optional CCD would
be turned off, and it is also possible that the second optional CCD would
be turned off.

The optimum ACIS-S spectroscopy chip set depends strongly on the expected
spectrum of the target. Typically the maximum signal is desired, so
the HETG and LETG observer is most likely to insist on all ACIS-S chips.
If the science does depend strongly on the flux received on the S0 and
S5 CCDs, the observer may need to specify all 6 ACIS-S CCDs to be on
(see Figure 6.41).
However, the observer should be aware that the amount of useful flux on
S0 and S5 with the LETG is typically quite low given they are both
FI CCDs (see Figure 9.6). Given the current
thermal performance of the spacecraft, it is probable that the first optional
CCD would be turned off and it is possible that the second
optional CCD would also be turned off.

We want to know the maximum number of counts in the spectrum from the brightest
source the proposer will be analyzing in the field (not necessarily the target),
and if the proposer believes that there will be lines in the spectrum
from this source. Both questions are being asked to trigger assessment
by the CXC of the sensitivity of the experiment to the gain
calibration and so that appropriate scheduling may
be employed to avoid thermally-induced gain drifts that might impact the
observation. If the GO is interested in the optimal spectral performance of the
ACIS CCDs, they should seriously consider using only 4 CCDs as
discussed above.

Optional Parameters that affect Pile-Up
See Section 6.12.1 for information on how the selection
of the number of operating CCDs and the size and position of the
subarray affects the minimum frame time. In general, using
fewer CCDs will permit a faster read-out. For example, using
the smallest subarray of 128 rows in the middle of the CCD has
a minimum frame time of 0.4s with 1 CCD operating and a minimum
frame time of 0.7s with 6 CCDs operating. In almost all cases,
the observer should leave the parameter on the RPS target
form that asks if the most efficient frame time should be
used at the default value of "Y". This parameter should
be changed only if the observer fully understands the impact on
efficiency of their observation. Please contact the
CXC HelpDesk (http://cxc.harvard.edu/helpdesk/) if you are considering a non-default setting
and are unsure of the CCD selection and CCD Frame exposure
time to use.

The Timed Exposure (TE) mode with the default nominal (and optimal) frame
time of 3.2s is the typical mode for ACIS observations. Note that the
option of selecting frame times shorter than nominal reduces observing
efficiency, and hence the number of photons collected for a given
observation time. (Note that the value of the nominal frame time
can differ a bit from 3.2s depending on how many CCDs are used; see
Section 6.12.1.)

Continuous Clocking Mode

The Continuous Clocking (CC) mode is useful when timing data are so
critical and/or pile-up is such a problem that the sacrifice of one
dimension of spatial data is warranted. The use of continuous clocking
may also lead one to consider specifying a particular satellite roll
orientation (see Chapter 3) to avoid having two
different sources produce events in the same CCD column. (See also
Section 6.21.4 below.)

This option applies only to Timed Exposure (TE) mode. The parameters
specifying an alternating exposure are:

the number of secondary exposures per primary exposure (1-15)

the primary exposure frame time

Frame times and efficiencies in TE mode are discussed in
Section 6.12.1, and the Alternating
Exposure option is discussed in Section 6.12.2.

Energy Filtering

It is possible to remove events from the telemetry stream, and thus
avoid telemetry saturation, by specifying an energy acceptance filter
within which detected events will be telemetered. The default discards
events above 3250 ADU (nominally 13 keV). The total per-chip
background rates for different upper energy cut-offs are in
Table 6.9.

Starting September 2006, a new energy to PHA conversion is used for
observations with energy filters. Two sets of conversions are
used, depending on the aimpoint of the observation. Observations with
ACIS-S at the aimpoint use a conversion tailored for the
BI CCDs and those with ACIS-I at the aimpoint use
FI CCD-specific conversions. The BI and FI specific
conversions are more accurate for each type of CCD than the conversion
used in previous cycles. The assumption is that it is desirable to have
the most accurate gain conversion for the CCD on which the
HRMA default aimpoint falls. Note that the conversion only impacts the on-orbit
energy filtering. Ground data processing
will always apply the appropriate PHA to energy conversion.

The observer should be aware that for observations which mix CCD types
(i.e. both BI and FI CCDs on), the selected
conversion (based on aimpoint
as above), will nevertheless apply to all selected CCDs. This will
not affect the
observation if the low energy threshold for the energy filter (the
"Event filter: Lower" parameter) is 0.5 keV or less, as the use of
either conversion at these energies results in essentially no
difference in the number of accepted events. However, for selection
of a low energy threshold above 0.5 keV, the
conversions are significantly different. Finally, observations which apply an
energy filter with a low energy threshold greater than 0.5 keV will
automatically be assigned spatial windows that allow the FI CCDs to
use the FI conversion and the BI CCDs to use the BI conversion
regardless of aimpoint. Proposers who need an energy filter lower limit above
0.5 keV are encouraged to contact the CXC HelpDesk
(http://cxc.harvard.edu/helpdesk/)
to discuss their
plans with an instrument scientist.

Spatial Windows

A more sophisticated approach to removing data from the telemetry
stream, and thus avoiding telemetry saturation, is by the use of a
Spatial Window. This option offers a good deal of flexibility. One may
define up to 6 Spatial Windows per CCD. Each window can be placed
anywhere on the chip. Note there is a significant difference between
a Spatial Window and a Subarray (Section 6.12.1):
Subarrays affect the transmission of CCD data to the on-board
ACIS processors; Spatial Windows select events detected by the processors
and only impact the telemetry rate. The user may also specify the
window energy threshold and energy range.

Spatial windows can specify the sample rate for events inside them.
A sample rate of 0 excludes all events; the default rate of 1
includes all events; a rate of n > 1 telemeters one out of every
n events in the window. A spatial window could be used to eliminate
a bright, off-axis source that would otherwise overwhelm the
telemetry stream. The order in which the spatial windows is
specified is important if they overlap. The first specified
window which includes a given pixel will be applied to events at that
pixel.

There are a small number of additional parameters that need to be
considered in specifying an observation with ACIS: (1) the off-axis
pointing (if required), which reduces the flux, and spreads out the
image; (2) the roll angle (Chapter 4); (3) time
constraints (if any); and (4) time monitoring intervals (if
any).

In the past, the Continuous Clocking (CC) mode
(see Section 6.12.3)
has been used to mitigate pile-up in very bright sources
(see Section 6.15.2).
ACIS has offered two standard telemetry formats for
observations performed in CC mode, "faint" and "graded"
(see Section 6.14.2).
While the CC faint mode choice has been used on occasion, the
CC graded mode has been employed for over 2 Msec
of Chandra observing time,
primarily to accommodate HETG spectra of bright X-ray binaries.
The objective was to use CC mode to mitigate
pile-up and conserve discrete structures such as emission and
absorption lines, and edges in the dispersed spectrum. CC graded mode also
reduces the possibility of telemetry saturation.

The calibrations of TE and CC modes are very similar and we have found
only small differences with respect to ACIS gain, response, and integrated
grade distribution, with CC faint mode versus CC graded mode choice
resulting in less than 3% differences at the most.
We have developed methods to assist the calibration in CC graded
mode in some rare cases. Starting in 2009-Nov, the CC faint
and graded modes have been altered to include some of the flight
grades that were previously rejected on-board: ACIS now telemeters
all flight grades except 24, 107, 127, 214, 223, 248, 251, 254, and 255.
While in faint mode these grades
are already recognized in the data processing, they will be accounted
for in the next CIAO release. It is recommended to use graded mode only
in extreme cases.

See Section 8.5.1 in Chapter 8 for
recommended ACIS modes, CCD choices, subarrays, etc.; that
section also discusses some aspects of analyzing the data.